1.1. Introduction IJ.\1F sector structure Solar wind at 1 AU Large scale features of the solar corona 19

Transcription

1 CHAPTER I HELIOSPHERIC CURRENT SHEET Al~D INTERPLANETARY l4edium page 1.1. Introduction IJ.\1F sector structure Solar wind at 1 AU Large scale features of the solar corona Heliospheric current sheet 28 I, i

2 1 CHAPTER I HELIOSPHERIC CURRENT SHEET AND INTERPLANETARY MEDIUM 1.1. Introduction The expanding solar corona carries the plasma and magnetic field of the solar atmosphere to great distances out in the solar system and provides the interplanetary plasma and magnetic field (IMF). The solar wind and IMP have extensively been studied during the past three decades. The interplanetary medium shows a large variety of short time scale and long time scale fluctuations depending on the phase of the solar activity cycle. The IMF is highly fluctuating on a short time scale, but on an average of few hours or more, it is ordered into large scale features. Sector structure is one such large scale feature, with average magnetic field directed either away from the sun (away sector) or toward the sun (toward sector) roughly along the Archimedean spiral direction defined by the zeroorder model of Parker (1958, 1963). The sector str~cture observed in the interplanetary medium (Ness and Wilcox, 1964; Wilcox and Ness, 1965) is now understood as a result of the presence of a heliospheric current sheet separating the heliosphere into two hemispheres of opposite dominant

3 2 magnetic polarity. The heliospheric current sheet is observed near earth as IMF sector boundary at which the I~1F field-direction reverses within few minutes to few hours. The number of sector boundaries observed near earth during a solar rotation period depends on the warpings of the HCS which in turn is determined by the complexity of the solar magnetic field. Thus the study of the morphology and periodic variations of the sector structure can give an idea of the complexity of the solar magnetic field. The study of the solar activity and its signature on the interplanetary medium is important in understanding the solar terrestrial relationship. In recent years, a wide variety of techniques and methods have been used to understand the solar processes and their relationship \vith the interplanetary medium. The aim of this chapter is to introduce some of the basic concepts available in literature on IMF sector structure, solar wind at 1 AU, large scale features of solar corona and heliospheric current sheet, so as to give a background to the problems discussed in the following chapters I~W sector structure The measurements of IMF polarity have been carried out using the magnetic field data obtained from earth orbiting satellites since 1964 (VJilcox, 1968; Couzens and

4 3 King, 1986). An elt8tnatp ~e~hod of infprring the of IMF near earth was discovered by Sva1gaard (1968) and Mansurov (1969) making use of the polar cap geomagnetic field observations. From this method, one can infer the I~lF polarities back to 1926 (Svalgaard, 1972; Mansurov and Manurova, 1970; Sva1gaard, 1976; Matsushita and Trotter, 1980). The inferred IMF polarity using ground based magnetograms of polar stations [Thule ( 'N) and Vostok ( 'S)J show a high degree of agreement with those determined using satellite observations and hence used for solar-terrestrial studies (Russel et al., 1975; Wilcox et al., 1975). Mori and Nagashima (1979) developed a method to infer IMP sector polarity from the cosmic ray north-south asymmetry, utilising ground level cosmic ray observations. Nefdner (1982) describes a method to infer IMF sector boundary passage from aa indices of geomagnetic activity after studying 72 disconnection events of 29 comets extending back to The mean field observations of solar magnetic field (magnetic field of sun as a star) is also another means to infer the IMF polarity near earth (Severny et ~l., 1970; Scherrer et al., 1977a,b). The IMF sector structure and sector boundary is related to a large number of solar-terrestrial phenomena. The IMF sector structure re~lect the large scale solar magnetic field (Svalgaard and Wilcox, 1975, 1976a,b; Wilcox

5 and Scherrer, 1981). So one can infer the large scale solar magnetic field by studying the properties of sector structure. The interplanetary plasma parameters like magnetic field intensity, particle density, temperature, solar wind velocity, cosmic ray flux, interplanetary ion flux and solar activity parameters like solar flare occurrence, calcium plage occurrence, green corona and K corona brightness, solar radio emission etc. are found to be organised around the IMF sector boundary (or RCS) in characteristic form (Wilcox, 1968, 1977). Apart from this, several terrestrial phenomena are also found to be organised around the IMF sector boundaries (Taylor, 1986; Svalgaard, 1977; Arora and Rangarajan, 1981). Heliolatitudinal variations in I~~ Most of the measurements of INF is linited to heliographic latitudes related to annual motion of earth around the Sun. There are also few out of ecliptic spacecraft measurements of INF. Rosenberg and Coleman (1969) proposed a dominant polarity effect of I~lF from analysis of IMF measurements from different space-crafts between 1962 and In their model, the dominant polarity of IMF in the northern or southern heliohemisphere coincides with the polarity of solar polar magnetic ~ield of the same solar

6 5 hemisphere (Rosenberg, 1970). This model also predicts a reversal of the dominant polarity of I~lF ln a particular heliohemisphere with the reversal of the solar polar field during the post sunspot maxima periods and the effect of 22 year heliomagnetic cycle on the IMF sector structure (Wilcox and Scherrer, 1972; Hedgecock, 1975; Smith et al., 1986). Latitude dependence of IMF polarity during is studied by Villan~e and Mariani (1983) using satellite observations. Eventhough, the dominant polarity effect proposed by Rosenberg and Coleman (1969) is apparent in near earth observations for few years, in general the statistical support is not satisfactory when long term IMF observations are analysed (~loussas and Tritakis, 1982; Tritakis, 1984a; Xanthakis et al.1981; Neidner,1982). Behannon et ale (1989 ) used source surface magnetic field data computed from potential field modelling of photospheric magnetic field observations for the period to show the predominant polarity variations of the IMF with heliographic latitude within Pioneer and Voyager missions provide out of ecliptic observations of IMF up to 27 0 north of heliographic equator and up to a radial distance of 25 AU from the Sun (Behannon et al., 1989; Mihalov et al., 1989; Smith et al., 1978).

7 6 Norphology of INF sectors Generally two or four IMF sectors are observed per solar rotation near 1 AU and occurrences of six sectors are rare. The complexity of the sector structure vary with solar sunspot cycle. T\vO sectors are more frequently observed in the high solar activity periods (Svalgaard et al., 1975; Sawyer, 1974) and four sectors are observed near minimum and post sunspot maxima periods. A typical large scale structure of IMF have life times up to a period greater than one year (Svalgaard et al., 1975). Behannon et ale (1989) obtained a mean life-time of 6 to 7 solar rotations for the heliospheric sector structure. The complexity of the sector structure vary with heliolatitude of observations (Sawyer, 1974; Behannon et al., 1989). The sector boundary also prefers to occur in certain heliolongitudes or certain days in Bartel's rotation period (Tritakis, 1979a, 1986). The number of sectors observed per solar rotation also depend on the relative contribution of various magnetic multipoles present in the solar magnetic field (Schultz, 1973; Levine, 1980a; Hoeksema, 1984). Wang Shui and Fang Lizhi (1979) discussed the IMF sector structure in terms of large scale spiral waves in the solar equatorial plane. Plyusnina (1985) observed a north-south asymmetry in the correlation between IMF and photospheric magnetic field structures in the

8 7 opposite heliohemispheres. Luhmann et ale (1987) reports on the asymmetry in the IMF strength about the heliographic equator from spacecraft observations in the interplanetary medium. Periodicities in the Il4F polarity Observations of IMF structure near lau generally show a 27 day recurrence period. Sometimes during solar maxima) a 28 day recurring IMF structure is also found to coexist with the 27 day pattern. The recurrent pattern of the IMF change with phase of the sunspot cycle and it remains relatively stable with a period. of 27 days during the declining phase of the solar cycle (Svalgaard 1972; Svalgaard et al., 1975). Hoeksema and Scherrer (1984) -- reports a also periodicity of 27 days and 28 days associated with equatorial dipole component of the coronal magnetic field. Sheeley and Devore (1986) and Nash et ale (1988) proposed a model to explain the day periods in IMF sector pattern. A comparison of coronal and IMF recurrence period during was reported by Parker (1987). Using inferred IMP polarities, between 1926 and 1986 Gonzalez and Gonzalez (1987) obtained various periodicities about 27.5, 13.5, 9.1 and 6.8 days associated with the evolution of IMF sector pattern. They also found periodicities of about 3.7

9 8 and 1.5 years associated with change ln the dominant structure of IMF from a well defined 2 sector to 4 sector structure. Xanthakis et ale (1981) also reports on the power spectral analysis of inferred IMF polarity ( ). Kotov and Levitskii (1984) found periodicities in the range days in the evolution of mean magnetic field of the Sun and IMF polarity from a power spectral analysis of solar and interplanetary data. Similar periodicities are also found to exist in geomagnetic and sunspot data (Gonzalez and Gonzalez, 1987; Villante and Francia, 1986, 1988) Solar wind at lau rrhe concept of coronal expansion out into interplanetary medium emerged as early as 1900 and gained support from the observations of comet tails (Vlilcox, 1976; Birkeland, 1909; Biermann, 1951). The spadecraft observations in early sixties confirmed the constant solar wind flow in the near earth environment. Basically there are types of solar wind flows observed near lad associated two with distinct solar origin and properties viz., low speed and high speed solar wind. This background flow of solar wind is occasionally modified by the transient flows from features in the Sun like solar flares or coronal mass ejections. Table 1.1 shows typical characteristics of three types of

10 9 Table 1.1 Comparison of solar wind parameters in three types of flow observed at 1 AU Parameter High speed solar ".,rind Lo\-v speed solar wind Solar wind relatedto solar transients v, km s n, em nv, em s nmv 2,dyne -2 em nv(mv 2 /2+rnHG/r), erg em -2 s T 105 K p' T 105 K e' n v/n v p p Ionisation temperatures 6 TO' 10 K? 6 T Fe, 10 K? Q 10 erg em s e' B, nt < up to 3.4 up to Field topology open? closed = 8 kt /B 2 1 p B /nv, 10-3 ntem 3 skm- 1 1 r

11 10 solar wind flows (Neugebauer, 1983). Low speed solar wind Low velocity solar wind with speed kms- l are generally associated with the closed magnetic field regions in the Sun with enhanced electron density viz., coronal streamers or heliospheric current sheet (Gosling et al., 1981; Esleev and Filippov, 1988). Solar wind properties show substantial variability, in most of the parameters in a low speed solar wind flow (Feldman et ~~., 1977). Proton temperature and momonentum flux density observed near lau in the low speed solar wind show an inverse relationship (Lopez and Freeman, 1986; Sastri, 1987). The helium abundance also observed to be lower in the low speed solar wind (Borrini et al., 1981; Feldman et al., 1981; Gosling et al., 1981). The slow speed solar wind carries away the positive angular momentum from the Sun and it is released from the Sun at significantly lower Alfven radius than the fast solar wind (Schwenn, 1987; pizzo et al., 1983). Kojima and Kakinuma (1987) reported the solar cycle variation in the lov! speed solar wind with respect to HCS using IPS (Interplanetary Scintillations) observations. Multi-spacecraft study of the spatial structure of solar wind velocity during is recently reported (r1iyake et al., 1988). Lyubimov and

12 11 Pereslegina (1985) discuss on the chromospheric and coronal sources of low speed solar wind. High speed solar wind Examination of solar wind data accumulated by different spacecrafts reveal continuous presence of various types of structures in the expanding solar wind. One of theq is the high speed stream which was identified in the Mariner 2 data of 1962 (Neugebauer and Snyder, 1966). A high speed plasma streaq is characterised by a large increase in the solar wind velocity lasting for several days. The flow tends to come from slightly east of the Sun as the speed begins to rise and form a westerly as the solar wind speed approaches its peak (Ness et al., 1971; Siscoe et al., 1969). Different definitions have been used to identify and catalogue high speed streams (Intrigilator, 1973, 1977; Bame et al., 1976; Gosling et al., 1976; Broussard et al., 1977; Lindblad and Lundstedt, 1982, 1983; Mavromichalaki, 1988). Generally, two different classes of high speed streams observed near 1 AU are (1) cororating streams flowing generally from coronal holes and (2) transient streams associated with solar flares (Iucci et al., 1979; Tsurutani et al., 1987). The basic features of cororating high speed streams are:

13 12 (a) The proton density (n) rises to unusually high values near the leading edge of the streams persisting for one day and generally having an inverse relationship to bulk speed (V). (b) The IHF magnitude (B) is proportional to bulk speed with constant IMF polarity throughout the stream except for fluctuations lasting for few hours. (c) Proton temperature (Tp) pattern is similar to bulk speed. The properties of flare-related stream are: (a) All the interplanetary parameters show simultaneous increase, possibly denoting radially outco~ing fast shocks. (b) V, Tp and B show large fluctuations in the maximum solar wind region. The polarity of the magnetic field often show inversions. (c) Tp behaviour tends to depart from solar wind behaviour. The high speed stream - IMF sector structure relation is investigated by many authors (Gosling et al., 1976, 1978; Hundhausen, 1977; Sawyer, 1976; Sheeley and Harvey, 1981; Sheeley et al., 1977; Obridko and Shelting, 1987a,b). The high speed steams of solar wind are generally found to be associated with single IMF polarity. There are

14 13 IMF sectors without high speed streams and sometimes more than one high speed stream are found within a IMF sector. The properties of high speed streams are found to evolve with solar cycle (Intrigilator, 1977; Bame et al., 1977). The preferred Bartel's day distribution of the occurrences of high speed streams during cycle 20 and 21 has been investigated (Lindblad, 1981; Rangarajan and Mavromicha1aki, 1989). Streams observed during the ascending phase and maxima of the sunspot cycle are generally short lived and non-recurring while stable recurrent high speed streams are observed during the declining period of sunspot cycle. The duration of high speed solar wind stream can vary from few days to nearly two weeks as observed near 1 AU. The amplitude of the high speed streams vary between kms. Broader streams are observed during low solar activity periods. Both equatorial coronal holes and equatorward extension of polar coronal holes are identified to be the source of high speed solar wind streams observed near the earth (Hundhausen, 1977). The amplitude or velocity maxima of high speed stearns are observed to be correlated with size, area and geometrical divergence of magnetic field lines of coronal holes (Nolte et al., 1976; Schwenn et al., 1978; Eslevich and Fi1ippov, 1986). The large scale magnetic field of the solar corona is observed to control the three dimensional density structure of the corona and associated

15 14 flow of solar wind into the interplanetary medium (Levine, 1978, 1980b; Levine et al., 1977; Pneumann, 1976; Pneumann et al., 1978; Burlaga~! al., 1978; Hundhausen, 1977). Shukhova et al. (1987) discusses on the open loops of IMF observed in the ecliptic high speed solar wind. Burlaga (1986) reviews our current knowledge on the structure and dynamics of corotating and transient solar wind streams in three dimensions. Henning et al. (1985) made a study on the observation of solar flares and related high speed streams near 1 AU in relation with the structure of the heliospheric current sheet. Apart from the observations of the solar wind near sun's equatorial plane by earth orbiting satellites IPS (Interplanetary scintillations) observations provide information regarding the solar wind flow at different heliolatitudes with in of the solar equator (Sime and Rickett, 1978, 1981; Coles et al., 1980; Rickett and Coles, 1983; Kojima and Kakinuma, 1987). Coles et al. (1980) found from IPS observations that the solar wind flow from the polar regions of the sun vary with sunspot cycle. The flow properties of the solar wind from the equatorial and polar region of the sun differ in many aspects),contributing to the observed solar wind velocity variations near 1 AU ~vithin ecliptic (Simon and Legrand, 1987). It is observed that K-corona brightness or electron density distribution in the the

16 15 solar corona is inversely correlated with solar wind speed (Sime and Rickett, 1978; Rickett and Coles, 1983). The variation of the properties of polar solar wind with sunspot activity is related to the corresponding change in the poloidal solar magnetic field and the associated change in flux tube divergence of magnetic field lines originating from polar coronal holes (Simon and Legrand, 1986; Coles al., 1980). Long term variations in the ecliptic solar et. wind can be inferred from geomagnetic or sunspot data (Legrand and Simon, 1981, 1985; Simon and Legrand, 1986, 1987; Russel, 1975; Sargent, 1986; Silverman, 1986). Despite considerable solar wind and related observations a~e accumulated over the past three decades the basic physics of the heating and acceleration of the solar wind is yet to be understood properly. Excellent reviews on this subject exists (Hundhausen, 1972; Hollweg, 1978; Withrobe, 1986; Holzer, 1979; Leer et al., 1982; Leer and Holzer, 1985; Pneumann, 1986; Marsch, 1985; Schwenn, 1987; Leer, 1987). Apart from physical mechanisms like purely thermal acceleration with and without extended heating, acceleration due to Alfven wave pressure, diamagnetic acceleration, the role of sporadic events such as spicules, macrospicules, X-ray bright points, ephemeral magnetic field regions and out flows seen in the EUV related to explosive events in the Sun are all explored in this connection

17 16 (Pneumann, 1986; Yang and Schunk 1989). Heliornagnetic latitudinal organisation of solar wind speed The concept that the three dimensional coronal magnetic field geometry controls the spatial structure of the solar wind in the heliosphere gained support through several studies after the skylab period (Levine et al., 1977; Hundhausen, 1977, 1978; Svalgaard and Wilcox, 1978; Schultz et al., 1978). Subsequently Hakamada and Akasofu (1981) demonstrated the feasibility of explaining most of the solar wind speed variations (daily, 27-day, semi-annual etc.) observed near 1 AU by assuming a positive gradient in solar wind speed with angular distance from the heliospheric current sheet (heliomagnetic latitude) and change in heliomagnetic latitude of the observing point (e.g. earth or space craft). Zhao and Hundhausen (1981) found a relation V(km/s) = sin2~ (1) during 1974 between solar wind speed V and the angular displacement ~ from a flat heliomagnetic equator inclined approximately deg with respect to the equatorial plane of the Sun. Zhao and Hundhausen (1983) found a new relationship V = * sin 2 p ( 2 )

18 17 for 1976 where p is the angular displacement from HCS inferred from K-corona observations and a plateau of 600 kms-1 for f3 > 135/day. Hakamada and Munakata (1984) obtained a relation V = sin 2 y (3) for the period wherey is the heliomagnetic latitude of the observing point derived from HCS inferences using potential field modelling of solar magnetic field. Coles et ale (1980) from a study of solar wind observations using IPS observations suggested that the heliobagnetic latitude dependence of solar wind speed can change with the sunspot cycle. This is confirmed by the study of Newkirk and Fisk (1985) concerning the heliomagnetic latitude dependence of solar wind speed using HCS inferenc8s from synoptic X-corona data for the period and subsequently by Fry and Akasofu (1987) investigating the sa~e using source surface magnetic field observations for the period It found that heliomagnetic latitudinal gradient of solar is wind speed is steeper during sunspot maxir;l.um. Bruno et ale 1986) investigated the latitudinal gradients of different solar wind parameters during using in situ observations. Recently Kotova et ale (1987) reported on heliomagnetic latitude dependence of solar wind speed using PROGN07.-9 measurements. One can tind evidence for correlati n ll between

19 18 solar wind speed and magnetic field strength in the source surface in various studies (Steinolfsen, 1982; Pneumann and Kopp, 1971; Yeh and Pneumann, 1977; Suess et al., 1977; Munro and Jackson, 1977; Hoeksema, 1984). Suess et ale (1984) found that the source surface field strength and solar wind speed near 1 AU correlated better than the heliomagnetic latitude during They also show the importance of removing effects due to transient related events when we study the heliomagnetic latitude dependence of solar wind speed near 1 AU. During 1974, Hakamada (1987) found that angular distance from the HCS is a better organiser of the solar wind speed than the source surface field strength. Kojima and Kakinuma (1987) investigated the solar cycle variations in the solar wind speed in relation with change in Hes structure using IPS observations and also suggested the limitations in expressing solar wind speed as a simple function~heliomagneticlatitude. Bruno et ale (1986) and Fry and Akasofu (1987) independently showed that solar wind speed distribution is asymmetric about the HCS. A 22 year cycle variation in the solar magnetic structure and solar wind is also proposed (Chirkov and Sansonov, 1984). According to Simon and Legrand (1987) the heliodagnetic latitude dependence of solar wind change from one solar cycle to another following the change in the maximum of sunspot activity.

20 Large scale features of the solar corona The white light corona becomes visible during total solar eclipses and there is an observational record of white light corona during total solar eclipses for nearly 100 years. Outside the eclipses it can be observed at 1-10 Re by balloon borne or satellite borne coronographs (e.g., 080-7, 8kylab) or ground based K-corona Qeters (e.g., Mauna Loa observations in Hawaii Islands). The K-corona or electron corona is produced by Thompson scattering of photospheric radiation by the electrons of the highly ionized coronal plasma. The K-corona is very inhomogeno~s containing a number of characteristic structures such as streamers, arches, plumes, fine rays and coronal holes. K-corona is observed by the Mauna Loa K corona meters since 1965 (Hundhausen et al., 1981). The data is presented in the form of synoptic charts, derived from successive daily scans in position angle around the solar limb. Each map is a cylindrical projection of coronal polarisation brightness (pb) plotted against solar latitude and longitude. The pb provides a measure of integrated electron density of the corona at a given height (1-4 R0) above the solar surface (Perry and Altschuler, 1973; Hundhausen, 1977). The K-corona pb map for the carrington rotation 1614 is shown in figure 1.1 (8ime, 1988).

21 20 - HIGH f,j..,tit\ ~ATCRi, HA~ LOA, ~11 CONiQRS (1:' K-C~ IPS) AT R ROT;HICN 1614, E:AST L1t13 OCGIt'NIN::; DAy 107.4, \974 ""1. t o I. I.. ~r"--._ a--i.-l-"""'-'+ -- _-I-1I_+-+-~---=:>.{-\,,..+---I 'I-.. ~ -"t' ; o ~.... "... - I 90 vo carrca..r:i ~ " 2, 3, 4, 6, 8, 10, 12 ffi X 10-8 Fig.l.l. K-corona brightness synoptic map for the carrington rotation 1614 (Sime, 1988)

22 21 Due to its high temperature of the order of million degrees, the solar corona shows line radiations from highly ionized atoms in emission against continuous spectrum. The most prominent lines are red (Fe X A 6374 " A) o 0 the green (Fe XIV ~ 5303 A) and the yellow (Ca XV ~ 5694 A) (Waldmeir, 1971). Outside the total solar eclipses these emission lines are observed employing a Lyot coronograph with appropriate narrow band filters of spectrograph. The o green corona (A 5303 A) has observational record since 1939 (Waldmeir, 1981; Leftus and Sykora, 1982). In addition to this, corona can also be observed in X-ray and XUV wave lengths using mainly space crafts (Broussard et al., 1978). o The absorption lines of Heliulu (He I 5876 A and He I o A) and radio observations can also give inforqation regarding the solar corona (Harvey and Sheeley, 1979). Coronal streamers Coronal stream.ev$ are approximately radial structures of high electron density found between 0.5 to 10 R. Active region steamers fording above young active regions are short lined, while helmet streamers lying above quiescent prominences or extended bipolar regions live for many months (Koutchmy, 1977). Streamers are generally associated with magnetic field structure that is closed and inhibit outflow of plasma (Pneuma~n et al., 1978). The

23 22 heliospheric current sheet which separates regions of opposite dominant magnetic polarity in the heliosphere is identified on the centre of the bright band of coronal streamers observable in K-corona (Hundhausen, 1977; Korzhov, 1977). Fisher and Sime have co~ducted a study on the coronal streamers using K-corona observations for the interval (quoted by Sime, 1986) using Mauna Loa K-corona observations. Number of streamers observed per solar rotation change with sunspot activity with maximum number being observed in sunspot maxinum and minimum number during sunspot minimum. Most of the structures occur near solar equator. The latitude of the highest streamer is In close correspondence with the position of HCS in comparison with other features like latitude of the highest filaments during a solar rotation period. The low speed solar wind flow ( km -1 s ) near 1 AU is associated with coronal streamers (Borrini et al., 1981; Feldman et al., 1981; Gosling et al., 1981; Sastry, 1987). Coronal holes Coronal holes are regions of abnormally low coronal density associated with magnetic field lines that are 'open' to interplanetary space. They were known to exist

24 24 Large coronal holes are present at the solar poles near declining phase and sunspot minimum and they shrunk in size during ascending phase of the solar cycle and disappear during the solar maxima. The area of the polar coronal holes waxes and and wanes with net amount of magnetic flux present at the solar poles (Broussard et al., 1978; Hundhausen e~ a~.) 1981; Webb and Davis, 1984). Coronal holes which occur at equatorial or middle solar latitudes are frequently connected to polar coronal holes, whenever sufficien~ unipolar regions develop at these sites. Their origin and evolution are more difficult to predict, but they are more stable during the declining phase of the sunspot cycle. Coronal holes are located within large scale magnetic areas dominated by one magnetic polarity and having a diverging (open) field geometry (Hundhausen, 1977). They for~ over regions where the magnetic field strength is locally high (0.7 G to 12 G) as shown in different studies (Bohlin and Sheeley, 1978; Harvey et al., 1982). A strong, large scale magnetic field will usually develop before a coronal hole and remain long after the coronal hole disappearance which suggest that coronal hole is an evolutionary step in the development of large scale magnetic regions in the Sun (Hoeksema, 1984). Coronal holes tend to occur in magnetic regions having the same polarity as the polar caps in the same

25 24 Large coronal holes are present at the solar poles near declining phase and sunspot minimum and they shrunk in size during ascending phase of the solar cycle and disappear during the solar maxima. The area of the polar coronal holes waxes and and wanes with net amount of magnetic flux present at the solar poles (Broussard et al., 1978; Hundhausen e/::. a.~') 1981; Webb and Davis, 1984). Coronal holes which occur at equatorial or middle solar latitudes are frequently connected to polar coronal holes, whenever sufficien~ unipolar regions develop at these sites. Their origin and evolution are more difficult to predict, but they are more stable during the declining phase of the sunspot cycle. Coronal holes are located within large scale magnetic areas dominated by one magnetic polarity and having a diverging (open) field geometry (Hundhausen, 1977). They for~ over regions where the magnetic field strength is locally high (0.7 G to 12 G) as shown in different studies (Bohlin and Sheeley, 1978; Harvey et al., 1982). A strong, large scale magnetic field will usually develop before a coronal hole and remain long after the coronal hole disappearance which suggest that coronal hole is an evolutionary step ln the development of large scale magnetic regions ln the Sun (Hoeksema, 1984). Coronal holes tend to occur in magnetic regions having the same polarity as the polar caps in the same

26 25 hemisphere but exceptions are observed (Harvey and Sheeley, 1979). Coronal holes exist directly adjacent to disk activity (Bohlin and Sheeley, 1978). The life time of coronal holes has a mean value of 6 solar rotations and range froq 3 to 20 solar rotations (Bohlin, 1977). The long term growth in size and decay of coronal holes occur at a rate of about 1.5 xl0 4 km 2 s- 1 which is consistent with the diffusion rate of magnetic field (Timothy et al., 1975; Bohlin, 1977). The area of coronal holes change mainly by sporadic, large scale shift of the boundaries (Nolte et al., 1978a,b). The boundaries of established coronal holes can be altered by near prominence eruptions and short lived coronal holes can be produced by these prominence eruptions and flares (vlebb et al., 1978). The coronal holes generally rotate rigidly with a synodic period of 27.2 days but occasionally some differential rotation is also observed (Timothy et al., 1975; Wagner, 1975; Bohlin, 1977; SheIke and Pande, 1985). The solar cycle evolution of coronal holes is explained in terms of locally flux iqbalance model of Broussard et ale (1978). There is a systenatic heliolatitudinal variation of occurrence of coronal holes during a solar cycle but indications of sone solar longitudinal organisation of coronal holes is also there (Broussard et eru.i') al., 1978; Hundhausen) 1981; Svalgaard and Duvall, 1977). The inter-relationship between coronal holes, solar wind

27 26 streams, IMF sector structure and geomagnetic activity during solar cycle 20 and early half of solar cycle 21 is given in a pictorial format by Sheeley et al. (1977) and Sheeley and Harvey (1981). Spatial and temporal changes in the coronal structure The overall shape of the corona is observed to change with the solar cycle. The ellipcity coefficient '~' characterising the shape of the isophotes of the white light corona is found to be small (~-0.05) during solar maximum and large ( -0.25) during solar minima (Koutchmy, 1977). The variation in the apparent shape of the corona arises partly from the changes in the distribution of bright features (steamers) and dark regions (coronal holes) observed around the limb in white light. Streamers are distributed fairly uniform around the limb of the Sun giving almost a circular shape of the corona during sunspot maxima. As the solar cycle progresses streamers tend to occur close to the solar equator and coronal holes establishes themselves at poles. During the declining phase and minimum of the sunspot cycle polar coronal holes dominate the high and mid latitudes of the Sun, while streamers are found only in low solar latitudes. During ascending phase streaders gradually spread

28 27 to higher solar latitudes and polar coronal holes \lill shrink and eventually disappear near solar maximauiundhausen et al., 1981). In the inner corona the magnetic field dominates the plasma \lhile in the outer corona radial flow of solar wind dominates. The coronal structure and its change reflects that of the solar magnetic field. At miniqa the large scale magnetic field is predodinantly dipolar and the open polar field lines provide large coronal holes at the poles. Long streamers extend far out along the equator marking the extension of surface fields by the solar \lind. At sunspot maximum the Qulti-polar components in the heliomagnetic fields are quite dominant and we see a bright corona all around the limb. Active regions provide coronal condensations and their magnetic remnants give rise to helmet streamers. Balough (1986) and f1c Queen (1986) reviews our current understanding on heliospheric and coronal magnetic fields respectively. The integrated brightness of the corona varies with sunspot activity as known from the study of Fisher and Bime (1984) using K-corona observations. In addition to the observed pole-equator asymmetry in the coronal brightness distribution, a north-south asymmetry in the same about the solar equator is also observed near sunspot minimum as evident from K-corona and green corona observations (Houssas

29 28 et al., 1982; Hundhausen et al., 1981). Tritakis et al. (1988) investigated the long term variations in the latitudinal and longitudinal asyn~etry of the solar coronal structure using green corona observations during Waldmeir (1981) studied the long term variations in area of the polar coronal holes bet~leen using green corona observations HELIOSPHERIC CURRENT SHEET Heliospheric Current Sheet (HCS) is a neutral sheet separating solar \vind flows carrying opposiuuy directed magnetic field in t~e interplanetary space. The idea of HCS was introduced by Schultz (1973) to explain the IMF sector structure and was developed later by Alfven (1977), Svalgaard et al. (1974, 1975), Svalgaard and Wilcox (1976b), Saito (1975) and Levy (1976). A phenoqenological model of HCS is developed by Kaburaki and Yoshii (1979) and Kaburaki and Imai (1983). A schematic drawing of the heliospheric current sheet is shown in figure 1.2.

30 29 +J OJ OJ..c: Ul +J l::: OJ H ::l o o.r-! H OJ..c: ~ Ul o.r-! rl OJ..c: OJ..c: +J 4-l o tj1 l:::.r-! ~ ru H '"d o.r-! +J ru Ei OJ..c: o U) N

31 30 Observations of heliospheric current sheet Potential field modeling Schatten et al. (1969) and Altschuler and Newkirk (1969) independently introduced the concept of a potential field model with a spherical source surface surrounding and concentric with -the Sun. In this model the line of sight photospheric magnetic field is used to determine the configuration of large Scale heliospheric magnetic field assuming that (1) There is no currents in the region bet\jeen Sun and the source surface and that (2) at the source surface the Qagnetic field is purely radial. The locus of the HCS is given by connecting the points where the radial field goes to zero at the source surface. Pneumann et al. (1978) and Wilcox et al. (1980) utilised this model to infer HCS structures during 1973 and 1976 respectively. Later Hoeksema (1984) refined this Qodel to coqpute HCS structure for the period using low resolution photospheric magnetic field observations made at the Hilcox Solar Observatory, Stanford. Several corrections

32 31 are to be applied before computing heliospheric magnetic field from line of sight photospheric magnet~c field. Some of them are (1) The polar field observed is to be corrected using annual variation in polar field strength (Svalgaard et al., 1978). (2) Zero-effect is to be removed from the data and (3) the source surface radii is fixed at an optiwum distance from the Sun (_2.5 R~) where the computed heliospheric field correlates better with the actual observations of IMF near earth. Further one needs to minimise the effects of evolution of large scale magnetic field within a solar rotation. Hoeksema and Scherrer (1986) extended the computation of heliospheric magnetic field up to 1985 and recently by Hoeksema (1989) up to Source surface magnetic field map for the carrington rotation 1740 is ShO\ln in figure 1.3.

33 ~ILCDX SOLAR 08SERUATORY SOURCE SURFACE r'lagnetic FIELD 0, ± 1,., '- 5, 10, 2 0 r"licrotesl':' 2 0 <: " I I-c ~±fi T::W~=e: o :) 0 18(1 210 <: <1 0 ",. (1 :;. (; (I (>.: ":, I~ I "7 d JI ",'./ Fig.l.3. Source surface magnetic (Hoeksema and Scherrer, field map 1986) for the carrington rotation 1740

34 33 K-corona observations The locus of the brightest regions in a synoptic K-corona map will give the position of HCS in heliocentric co-ordinates (Hundhausen, 1977). This method, known as maximum brightness curve (HBC) method is utilised by Burlaga et al. (1981) and Bruno et al (1982, 1984) to infer HCS geometry during Bruno et al. (1984) found that the RCS structures cornputed from K-corona observations show generally good agreement \lith the RCS inferred fror:1 potential field modeling a photospheric magnetic field data during 1976 except for minor differences. This result can also be found from similar other studies (Pneurnann, 1976i Pneumann et al., 1978i Wilcox and Rundhausen, 1983). -- Korzhov (1977) has used the OSO-7 K-corona observations and synoptic K-corona observations from ground based K-corona meters to infer RCS structures during (Korzhov, 1982). Newkirk and Fisk (1985) identifies the RCS on the centre of the bright band of coronal streamers in the K-corona observations for the period Saito and Swinson (1986) utilised the mid-line method to infer HCS position from K-corona observations. In their method RCS latitude at a given longitude is given by the mid-points of the boundary of polar coronal holes in the opposite

35 34 heliohemispheres. All the above investigations found a good correlation between the predicted polarity of IMF from their inferred HCS positions and actual observations of IMF near 1 AU. The inclination of HCS observed in K-corona is compared with the same observed near earth's orbit by Behannon et ale (1983). Other methods Yearly average solar wind maps covering nearly of solar latitude and entire solar longitude can be made using IPS observations (Coles et al., 1980 i Rickett and Coles, 1983; Kojima ar:d Kakinuma, 1987). The position ot ImJ speed solar wind belt (V<400-1 kms ) in these maps give a good approximation to the HC~ POS1L.LUIlS inferred using K- corona odservations or potential field modeling of solar magnetic field. Stewart (1985, 1987) developed a method to infer RCS position using synoptic plots of solar radio noise storms in the interval The dividing line between opposite noise-storm polarities appears to be a good representation of the HCS up to latitudinal displacements from the solar equator. In addition to these methods IMF p01arity observed near earth can also be used to infer the properties of HCS

36 35 (Svalgaard and Wilcox, 1976b; Tritakis, 1984a; Villante et al., 1979). Geomagnetic indices also provide clues to the structure of the HCS (Simon and Legrand, 1987; Saito and Saito, 1986; Triskova, 1988; Oksman and Kataja, 1986). The interaction between comet tails and the HCS is also used to infer HCS properties in the past (Neidner, 1982). Evolution of HCS with sunspot cycle During sunspot minimum the HCS stays very close to the solar equator. For e.g. HCS is within during several solar rotations in 1976 (Hoeksema, 1984; Newkirk and Fisk, 1985). But as the solar activity builds up, HCS extends to higher solar latitudes typically >70 0. Even though the structure of IMF sector structure observed near earth is simple consisting of either two or four sectors per solar rotation. HCS configuration during sunspot maximum is complex. Multiple HCS structures are usually observed during this period (Hoeksema, 1984; Korzhov, 1977). During solar cycle 21 through 1978 to 1981, the latitudinal extension of HCS was >50 0 But later after 1985, the latitudinal extension dropped to and further decreased near 1987 (Behannon et al., 1988). Thus one can find a systematic variation of the latitudinal extension of the HCS with the solar cycle. In addition to this, the number of warps in the

37 36 HCS also vary with sunspot cycle. In solar cycle 21 the warps of HCS extended up to 50 0 during several years of observation (Hoeksema and Scherrer, 1986). The geometry of HCS is not observed to be symmetric about the heliographic equator. During the HCS is extended more in the southern heliosphere (Burlaga et al., 1981; Korzhov, 1983). Tritakis (1984a) provides evidence of the asymmetry in HCS placement about solar equator from long term analysis" of IMF polarity data. The complexity of IMF sector structure observed near earth (two sector structure or four sector structure) depend on the structure of the HCS during various phases of the sunspot cycle (Hoeksema, 1984; Gonzalez and Gonzalez, 1987). Hundhausen (1977) and Thomas and Smith (1981) assumed a tilted dipole heliomagnetic configuration to explain the structure of the HCS. But during most of the solar cycle period higher order solar magnetic multipoles distort the structure of HCS from a simple sinusoidal structure due to a dipolar solar magnetic field. Even near sunspot minimum when the heliomagnetic field is predominantly dipolar, the structure of the HCS in the equatorial region of the Sun depends on a smaller quadrupolar field (Bruno et al., 1982). During most of the period one can approximate the large scale heliomagnetic field with a magnetic dipole and a magnetic quadrupole

38 37 contributing to the structure of the HCS (Hoeksema, 1984; Saito and Swinson, 1986; Levine, 1980). The contribution of higher order multipoles to the heliomagnetic field increases with sunspot activity. During high solar active period often the quadrupolar contribution exceed the dipolar contributions to the HCS and the octupole contribution becomes comparable to that of dipole (Hoeksema, 1984; Balough, 1986). Near solar maximum the HCS configuration is very complex and multiple HCS structures are commonly observed during this period. Another important feature is that the dominant polarity of IMP above or below the HCS is controlled by the 22 year heliomagnetic cycle. During solar cycle 21 the dominant polarity above or below the HCS reversed with the reversal ln the polarity of solar poles during 1980 (Hoeksema, 1984; Smith et al., 1986; Saito et al., 1987; Behannon et al., 1989). Microstructure of the HCS It was found from the space-craft observations of IMP sector boundaries near 1 AU that HCS is~ non-null magnetic field region where interesting reconnective processes may take place (Klein and Burlaga, 1980; Behannon and Neubauer, 1981; Behannon et al., 1981; Villante and

39 38 Bruno, 1982). Behannon et ale (1981) conducted a detailed study on the fine scale characteristics of IMF sector boundaries observed near earth during The field configuration of the HCS is found to be associated with directional discontinuitiesj 'thick' as well as portions of the HCS is observed. Behannon et ale 'thin' (1981) observed a thickness of 3xl0 14 cm associated with the HCS. Eslevich and Filippov (1988) show that the plasma and magnetic field in the HCS does not undergo appreciable additional change due to.interplanetary dynamics during its transit from Sun to 1 AU. HCS structure observed at the earths orbit is then basically determined by its projections at the solar surface. The variation of magnetic field in the HCS implies a rotation or tangential transition which the magnetic vector experience as they move across the HCS titled at angle 8 with respect to the ecliptic plane. The rotation of magnetic vector inside the HCS is likely due to the differential rotation of the Sun. Bruno and Bavassano (1987) investigated on the dependence of IMF power spectrum upon the angular displacement from HCS using Helios observations. At the low wave numbers (K<2xlO- 6 km- l ) the magnetic field controlled by slow solar wind has spectral characteristics different from those of fast wind. Close to the HCS a -1 spectrum as K is seen. Recently Bruno and Bavassano (1988) studied the variation of the power of IMF

40 39 fluctuations \vith angular distance from heliospheric current sheet. Deformation of HCS due to Solar wind velocity gradients In the highly idealised case of a totally uniform, steady, solar wind, the shape of the HCS is independent of the distance from the Sun. Suess and Hildner (1985) have considered the effect of an inhomogenous solar wind with velocity varying from point to point along the HCS near the source surface causing a distortion in the HCS geometry at large distances from the Sun. Suess et al. (1986) demonstrated this effect from observed azimuthal velocity gradient along the Hes during a solar rotation period in But the observations of the distant HCS structure by Pioneer and Voyager missions does not show appreciable distortions from the near earth HCS structure as predicted by Suess and Hildner (1985) (Behannon et al., 1989; Mihalov et al., 1989). Relation with cosmic ray propagation and comets Svalgaard and Wilcox (1976b) suggested that the solar cycle variations on the cosmic ray intensity observed near earth could be explained in terms of the change in tilt

41 40 of the HCS with solar cycle. Jokipii and Thomas (1981) were able to model the effect of simple tilted dipole configuration of HCS on the propagation of galactic cosmic rays. Heliospheric current sheet is thus a natural defence against intense galactic cosmic rays entering the heliosphere. Increasing the tilt of the HCS significantly decreases the flux of the cosmic rays at earth. During solar cycle 21 and near the minima of solar cycle 22 the changes in the tilt of the HCS fairly correlated with the cosmic ray intensity variations observed at earth (vvebber and Lockwood, 1988; Hoeksema, 1984). Neidner and Brandt (1978) suggested that the observed disconnection events (DE) of the ionic tails of the comets from the nucleus is associated with the interaction of the cometary plasma with the HCS. Neidner (1982) identified such 72 DE from observations of comet tails since 1892 and he has inferred the latitudinal extent and tilt properties of the HCS during these events. During the recent spacecraft encounters with comet Halley several DE were observed in relation with the HCS (Brandt and Neidner, 1987; Saito et al., 1987; Lundstedt and Magnusson, 1987).